DOI: 10.2298/SAJ1081001J Invited review
ASTRONOMICAL OPTICAL INTERFEROMETRY.
I. METHODS AND INSTRUMENTATION
S. Jankov
Astronomical Observatory Belgrade, Volgina 7, 11060 Belgrade 38, Serbia
E–mail: [email protected]
(Received: November 17, 2010; Accepted: November 17, 2010)
SUMMARY: Previous decade has seen an achievement of large interferometric projects including 8-10m telescopes and 100m class baselines. Modern computer and control technology has enabled the interferometric combination of light from separate telescopes also in the visible and infrared regimes. Imaging with milli-arcsecond (mas) resolution and astrometry with micro-milli-arcsecond (µas) precision have thus become reality. Here, I review the methods and instrumentation cor-responding to the current state in the field of astronomical optical interferometry. First, this review summarizes the development from the pioneering works of Fizeau and Michelson. Next, the fundamental observables are described, followed by the discussion of the basic design principles of modern interferometers. The basic inter-ferometric techniques such as speckle and aperture masking interferometry, aperture synthesis and nulling interferometry are disscused as well. Using the experience of past and existing facilities to illustrate important points, I consider particularly the new generation of large interferometers that has been recently commissioned (most notably, the CHARA, Keck, VLT and LBT Interferometers). Finally, I discuss the longer-term future of optical interferometry, including the possibilities of new large-scale ground-based projects and prospects for space interferometry.
Key words. Instrumentation: interferometers – Methods: observational – Tech-niques: high angular resolution
1. INTRODUCTION
An optical (visible and infrared) long-baseline interferometer is a device that allows astronomers to achieve a higher angular resolution than it is possi-ble with conventional telescopes. In fact, at wave-length λ, the resolution of a single telescope with aperture diameter D scales as D/λ while the reso-lution of two-telescope interferometer scales asB/λ, where the baselineBis the distance between the tele-scopes. However, more severely, both are limited by the atmospheric turbulence (seeing typically≈1 arc-second). For this reason, a single aperture telescope
Earth rotates. In addition, most of the new inter-ferometers have more than two mirrors in the array and can move the apertures along tracks.
The implementation of interferometry in opti-cal astronomy began more than a century ago with the work of Fizeau, who outlined basic concept of stellar interferometry, how interference of light can be used to measure the sizes of stars. He suggested that one could observe stars using a mask with two holes in front of a large telescope and that ”it should be possible to obtain some information on the angu-lar diameters of these stars” (Fizeau 1868). The first attempts to apply this technique were carried out soon thereafter (St´ephan 1874), although the largest reflecting telescope in existence at the time, the 80cm reflector at the Observatoire de Marseille, could not resolve any stars with a mask that had two holes sep-arated by 65cm, and only the upper limit (0.158 arc-sec) of stellar diameter could be derived. However, with a Fizeau interferometer on the 1m Yerkes refrac-tor, Michelson (1891a,b) measured the diameters of Jupiter’s Galilean satellites. Following Michelson’s success Schwarzschild (1896) managed to resolve a number of double stars with grating interferometer on 25cm Munich Observatory telescope, while not long afterward, on the 2.5m telescope on Mt. Wilson, Anderson (1920) determined the angular separation of spectroscopic binary star Capella (αAur).
Although the first measurement of a stellar an-gular diameter was performed on the supergiant star Betelgeuse (α Orionis) with Michelson’s 6m stellar interferometer in 1920 (Michelson and Pease 1921), the optical interferometry was slowly evolving from a difficult laboratory experiment to a mainstream ob-servational technique. Following the success of the 6m interferometer, Pease (with Hale) constructed a 15m interferometer but this experiment was not very successful. Due to the disappointing results from the 15m interferometer, it would be decades before sig-nificant developments inspired new activity in the optical field. The real difficulty is to combine the beams in phase with each other after they have tra-versed exactly the same optical path from the source through the atmosphere, each telescope, and further to the beam recombination point. This has to be done to an accuracy of a few tenths of the wave-length, which in the case of visible light, is not an ob-vious task, particularly because of atmospheric tur-bulence which makes the apparent position of a star on the sky jitter irregularly. This jitter often causes the beams in different arms of the interferometer to overlap imperfectly or not at all at any given mo-ment. For these reasons, the optical interferome-try requires extreme mechanical stability, sensitive detectors with good time resolution, and at least a simple adaptive optical system to reduce the effects of atmospheric turbulence. For these reasons, only the technologies emerging at the end of 20th century allowed the full application of optical interferometry in astronomy.
Meanwhile, advances in radar technique dur-ing World War II stimulated rapid development of radio interferometry beginning with the first radio interferometer built by Ryle and Vonberg (1946). The main reason for tremendous success
concern-ing astronomical applications is that the effects of atmospheric noise are much less pronounced at ra-dio wavelengths. Paralelly, the development of In-tensity interferometry discovered by Hanbury Brown and Twiss (1956a) inspired a new project in optical interferometry. Their basic principle describes how correlations of intensities (not electric fields) can be used to measure stellar diameters. The important point is that the technique relies on the correlation between the (relatively) low-frequency intensity fluc-tuations at different detectors, and that it does not rely on the relative phase of optical waves at the different detectors. The requirements for the me-chanical and optical tolerances of an intensity in-terferometer are therefore much less stringent than in the case of ”direct detection” schemes as it is Fizeau/Michelson interferometery. First results were reported soon thereafter (Hanbury Brown and Twiss 1956b), leading to the development of the Narrabri intensity interferometer. With a 188m longest base-line and blue-sensitivity, this project had a profound impact on the field of optical interferometry, mea-suring dozens of hot-star diameters (e.g. Hanbury Brown et al. 1967a,b, 1970, 1974a,b, Davis et al. 1970). The small bandwidths attainable with in-tensity interferometry limited the technique to the brightest stars, and pushed the development of so-called ”direct detection” schemes, where the light is combined before detection to allow large observing bandwidths.
With the advent of lasers in visible and in-frared, the benefits of radio interferometry have been pursued in optical long-baseline interferometry through heterodyne detection (e.g. Gay and Journet 1973, Assus et al. 1979). Radio interferometry func-tions in a fundamentally different way from optical interferometry. Radio telescope arrays are hetero-dyne, meaning that incoming radiation is interfered with a local oscillator signal before detection. The signal can then be amplified and correlated with sig-nals from other telescopes to extract visibility mea-surements. Optical interferometers are traditionally homodyne, meaning that incoming radiation is in-terfered only with light from other telescope. This requires transport of the light to a central station, without the benefit of being able to amplify the sig-nal. The other important benefit of radio interfer-ometry is that the required accuracy for beam re-combination is much more easy to achieve in radio then in optical domain. One heterodyne optical in-terferometer (ISI, see Section 3.1) has been built to operate at≈10µm wavelengths. As for Intensity in-terferometry, the technique is feasible but with small bandwidths attainable and limited to bright sources, while the homodyne optical interferometry allows large bandwidths to be used since the interfered light is detected directly. Two important steps towards modern optical interferometry have been done at the Observatoire de Calern, France:
free of turbulent areas, in essence an image of the object as it might appear free of atmospheric turbu-lence.
2. Long-baseline interferometry on a star by directly combining the electric fields before photon detection (”direct detection”) was first accomplished in 1974 by Labeyrie (1975) by using a 12m baseline at the I2T interferometer.
The following section deals with the funda-mental observables in astronomical optical interfer-ometry. In Section 3. three basic designs to achieve coherent combination of light are explained in more details, while in Section 4. the basic interferometric techniques are presented. Further development in the field is described in Sections 5. and 6. while the possibilities of new ground-based and space projects are discussed in Section 7. Finally, I summarize the most important conclusions in Section 8.
2. INTERFEROMETRIC OBSERVABLES 2.1. Visibility modulus
The fringe contrast observed with an interfer-ometer is historically called the visibility and, for the simple (two-slit) interferometer can be written as:
V =Imax−Imin Imax+Imin =
Fringe amplitude Average intensity (1)
where Imax andImin denote the maximum and min-imum intensity of the fringes.
Generally, the primary observable of an inter-ferometer is the complex visibility function (the am-plitude and phase), which is the Fourier transform of the object brigthness distribution (van Cittert-Zernike Theorem). Let ˜I(u, v, λ) denote the spa-tial Fourier transform of sky brightness I(s, p, λ) at the wavelength λ, where the Fourier frequency pair u, v (spatial frequency) is the conjugate of the spa-tial Cartesian coordinates s, p, respectively. Defin-ing the complex visibility ˜V(u, v, λ) = ˜I(u, v, λ)/N where N is normalization factor defined so that |V˜(0,0, λ)|= 1,it follows:
V =|V˜(u, v, λ)|.
The visibility V as defined by Eq. (1) is exactly proportional to the amplitude modulus of the image Fourier component corresponding to the spatial fre-quency (u, v) =~b/λwhere~bis the baseline vectorB~ projected onto the plane of the sky.
The phase Φ of the fringe pattern is equal to the Fourier phase of the same spatial frequency com-ponent but it cannot be observed since an incoming plane wave from a source is corrupted as it propa-gates through the turbulent atmosphere which pro-duces atmospheric phase delays such that the phase information on the source is completely lost. Prac-tically only the modulus of the complex visibility V =|V˜(u, v, λ)|(or squared visibility amplitudeV2) is an observable quantity.
The corruption of this phase information has serious consequences, since imaging of non-centrosymmetric objects rely on the Fourier phase information encoded in this intrinsic phase of inter-ferometer fringes. Without this information, imag-ing cannot be done except for simple objects such as discs or round stars. A number of strategies have been developed to circumvent these difficulties.
2.2. Relative phase
The technique called phase referencing relies on the fact that the relative phase can be obtained if atmospheric and instrumental noise can be con-trolled or compensated, for example through ref-erencing by measuring the instantaneous phases of fringes using a reference star (Shao and Colavita 1992) or by using the same source at another wave-length. The technique, called ”Differential Speckle Interferometry” (DSI) measures the relative shift of an object in different spectral bands, and its appli-cation is based on the assumption that the shift be-tween two speckle images (instantaneous distorted images, see Section 4.1) can be related to spatial structures of the object under study. Although the phase of the spatial coherence function is cor-rupted by instrumental and atmospheric noise, dif-ferent methods of data processing can be used to overcome this difficulty. If the two sets of speckle interferograms of the same object are recorded si-multaneously at two different wavelengths, but close enough for the Point Spread Function to be the same for both channels, then their ensemble average cross-spectrum provides the information about the rela-tive shift as a fraction of the object size. This tech-nique was described by Beckers (1982), Jankov et al. (2001).
2.3. Closure phases
The closure phase measurements tell us about the symmetry of an object and is only available from a closed triangle of elements, such as three tele-scopes. Normally, atmospheric fluctuations disturb the fringe phase between any two interferometer el-ements, but these fluctuations cancel out in a closed triangle.
Considering a 3-telescope array where the phases Φij, Φjk, Φkiare intrinsic phases, Φ′
ij, Φ′jk, Φ′ki are phases measured on the baselines between tele-scopes (i and j), (j and k), (k and i) respectively,ψi, ψj, ψk are the atmospheric errors introduced above telescopes (i,j,k) andεij, εjk, εki is the noise, it fol-lows:
Φ′
ij= Φij+ψi−ψj+εij
Φ′
jk= Φjk+ψj−ψk+εjk
Φ′
ki= Φki+ψk−ψi+εki
Consequently the closure phase:
Φcl= Φ′
ij+ Φ′jk+ Φ′ki= Φij+ Φjk+ Φki+εij+εjk+εki
2.4. Bispectrum
When analysing the speckle interferometry data, one can think of many virtual sub-apertures (with sizes equal to the coherence length λ¯2
∆λ) spread across the full telescope, with fringes forming be-tween all the sub-aperture pairs. The speckle in-terferometry data can be reduced using the ”bispec-trum”, which permits a direct inversion from the es-timated Fourier amplitudes and phases. The bispec-trum:
˜
Bijk= ˜VijVjk˜ Vki˜
is formed through triple products of the complex vis-ibilities around a closed triangle, where i,j,k specifies the three aperture locations on the pupil of the tele-scope.
It was discovered that the Fourier phases could also be estimated from such data (e.g. Knox and Thompson 1974), and that the bispectrum phase is identical to the closure phase. Interest-ingly, the use of the bispectrum for reconstruct-ing diffraction-limited images was developed inde-pendently (Weigelt 1977) of the closure phase tech-niques, and the connection between the approaches was realized only later (Roddier 1986, Cornwell 1987).
3. INSTRUMENTAL DESIGNS 3.1. Heterodyne interferometer
The principle of a Heterodyne interferometer is to change the frequency of light at each telescope, carry to the common focus an intermediate frequency signal, and combine all these signals in order to pro-duce fringes. Associated to each telescope is a laser, and the two beams (laser and star) are then com-bined with a beam splitter onto an infrared-sensitive photodiode, which produces an electric current pro-portional to the light intensity on it. Two light waves, those from the laser and from the star, are very close in frequency and while combining two beams, they beat together to produce a signal at lower frequency such that it becomes a radio wave.
To produce interference fringes two radio sig-nals are overlapped instead of light beams. All of the information that is carried in the visible or in-frared light wave from the star is still in the new radio signal, the only thing that is different is the wavelength. The principle of heterodyne interferom-etry is depicted in Fig. 1.
This method of converting the frequency of a signal from the star by mixing it with a fixed-frequency ”local oscillator” is called heterodyne de-tection. (The term ”local oscillator” refers to the laser in this case). Usually, the local oscillators are CO2 lasers operating at frequencies ≈ 10µm wavelengths. The heterodyne detection suffers from bandwidth limitations as well as additional noise contribution from laser shot-noise, which becomes progressively worse at higher frequencies. However, one of the main advantages is that the signals are converted to radio frequencies and then combined electronically using techniques from radio astronomy.
Fig. 1. Heterodyne interferometer arrangement. The light from an infrared laser together with infrared light from a star are converted to radio frequencies through heterodyne circuits and then combined elec-tronically. A light wave from a star hits first tele-scope, and has to travel an extra distance (d) before it arrives at second telescope. In order for the in-terference method to work properly, the signal from one telescope should be delayed in time (τ) until the wavefront reaches the other telescope; only then two signals can interfere together.
3.2. Intensity interferometer
The Intensity interferometer is also called ”the post-detection correlation Interferometer” since the signals are correlated after detection. The inten-sity interferometry technique does not rely on the actual interference of light rays, thus no fringes are formed. Instead, the interferometric signal, the gree of mutual coherence, is characterized by the de-gree of correlation of light intensity fluctuations ob-served at two different detectors. In practice, the de-gree of mutual coherence is measured using fast tem-poral correlations between narrow optical waveband intensity fluctuations observed by two (or more) tele-scopes separated by a baseline distance (Fig. 2).
The principal component of the intensity fluc-tuation is the classical shot noise which will not demonstrate any correlation between the two sepa-rated telescopes. The intensity interferometric signal is related to a smaller noise component: the wave noise. The wave noise can be understood as the ”beat frequency” in optical intensity between the dif-ferent Fourier components of the light reaching the telescopes. This wave noise will show correlation be-tween the two detectors, provided there is some de-gree of mutual coherence between the light received at the two telescopes.
The technique has the benefit of being much less sensitive to atmospheric phase fluctuations be-cause the signals input to the correlator have followed roughly the same path through the atmosphere. De-lay tracking and local oscillator stability do not have to be very accurate. Unlike the direct (Amplitude interferometry) the Intensity interferometry is well matched to blue/visible operation, and scales to long baselines with ease (the atmospheric stability re-quirements are a factor of a million less stringent).
The major drawback of Intensity Interferom-etry is that copious quantities of light are necessary to observe the wave noise signal in the presence of the much larger shot noise. Even for brightest stars, very large light collectors (≈ 100m2 area) are nec-essary to have a sufficient statistical strength to ob-serve non-Poisson fluctuations around the mean pho-ton intensity. Mainly for this reason, the experiment conducted at Narrabri from 1963 to 1971 was the only implementation of a stellar intensity interfer-ometer. Since then, all stellar interferometers have been based on Fizeau/Michelson’s techniques.
3.3. Direct interferometer
In this approach the light from each aperture is carried to a common focus, combined coherently in order to detect interferometric signal (fringes). The combination of two (or more) light beams can be done in the image plane (Fizeau mode, Fig. 3a) or in the pupil plane (Michelson mode, Fig. 3b). In the Fizeau mode, the ratio of aperture diame-ter/separation is constant from light collection to re-combination in the image-plane (homothetic pupil). In the Michelson mode, this ratio is not constant since the collimated beams have the same diameter from the output of the telescope to the recombination lens. This means the object-image relationship can no longer be described as a convolution. However, it is possible to use a relay optics after the beam-combiner so as to stack two output pupils on the top of each other with a modulation depending on the output Airy discs distance (Fig. 3c). The main ad-vantage over the Michelson mode recombinations is conservation of the object-image convolution relation (Vakili et al. 2004).
Interferometers in which the output pupil is not a homothetic image of the input pupil, but the pattern of the subaperture centers is conserved are known as densified pupil interferometers (Labeyrie 1996). In such an arrangement, the interferometer
forms a true image of small sources at the combined focus. This is particularly attractive for arrays con-sisting of many telescopes, because the desired infor-mation can be obtained after deconvolution with the known point spread function.
Fig. 3. Comparison of three different beam-combinations for an optical stellar interferometer. (A) and (B) the classical Fizeau versus Michelson beam-combinations, (C) instead of superimposing the Airy patterns from the telescopes, it is possible to use a relay lens after the beam-combiner so as to stack the two output pupils on the top of each other with a modulation depending on the output Airy discs dis-tance.
The disadvantage of the Michelson mode is a very narrow field of view compared to the Fizeau mode, but for diluted optical arrays it is extremely difficult to construct an interferometer with image plane beam recombination. For this reason Fizeau mode is more typically implemented on a common-mount interferometers. Typical example of such an interferometer designed to create high resolution im-ages over a wide field of view is the Large Binocu-lar Telescope (see Section 6.7), a single-mount struc-ture resulting in a fixed entrance pupil for which the Fizeau mode recombination is adapted.
The Michelson mode interferometry is used in long-baseline interferometers made of independent telescopes. Beam transfer optics and delay lines al-low it to maintain equal path length between the telescopes as the object is tracked across the sky.
Fig. 4. Michelson’s interferometer. The light from each aperture is carried to a common focus, combined coherently in order to detect interferometric signal (fringes). Beam transfer optics and delay lines al-low it to maintain equal path length between the tele-scopes as the object is tracked across the sky.
In addition, mechanical constraints on the in-strument, errors on the pointing model, thermal drifts, various vibrations and atmospheric turbulence make the null Optical Path Difference (OPD) point changing. The error on the OPD must be less than the coherence length defined by: lc =
¯ λ2
∆λ where ¯
λ is the mean wavelength observed and ∆λ is the spectral interval. Fringes are searched by adjusting the delay-line position, and this real-time control is called ”fringetracking”.
After being collected by the telescopes and properly delayed, the light must be directed to a central facility for beam combination. The interfer-ogram can then be recorded, as long as the scanning takes place faster than the atmospheric coherence time τc=c/∆λ.
4. INTERFEROMETRIC TECHNIQUES
4.1. Speckle Interferometry
Speckle interferometry is a method permit-ting the extraction of spatial information from two-dimensional images at scales down to the diffraction limit of the telescope, in spite of severe blurring in-troduced by atmospheric turbulence. The image of a star obtained through a large telescope looks ”speck-led” or grainy because different parts of the image are blurred by small areas of turbulence in the earth’s atmosphere. The image obtained for an unresolved point source depends on the exposure time. For long exposures, the image is blurred to a seeing disk ≈1 arcsec. For exposures shorter than the coherence time (tc ≈10ms for the atmosphere in the optical; 100ms for the infrared), a group of bright speckles is obtained approximately of the size of the Airy disk. In that way the speckle interferometry freezes the atmosphere, and gets an instantaneously distorted image as presented in Fig. 5.
Fig. 5. Speckle, an instantaneously distorted stellar image.
In 1970, the French astronomer Antoine Labeyrie showed that information could be ob-tained on the high-resolution structure of the object from the speckle patterns using the Fourier analysis (”speckle interferometry”). In the 1980s, the meth-ods which allowed images to be reconstructed inter-ferometrically from these speckle patterns were de-veloped using a sequence of short-exposure snapshots to obtain images at a telescope’s diffraction limit.
With existing large telescopes, speckle tech-niques thus permit resolution at spatial scales of 0.025 arcseconds rather than the 1 to 2 arcseconds associated with classical techniques. These meth-ods are also characterized by enhanced measure-ment accuracy of the separation of closely spaced objects seen through the turbulent atmosphere. The speckle interferometry system incorporates an inten-sified charge coupled device array as the primary imaging detector and an autocorrelator as a high speed data reduction processor operating at video rates. Also, the reduction of the speckle patterns can be done a posteriori involving computer processing.
4.2. Aperture Masking Interferometry
Aperture Masking Interferometry is a form of speckle interferometry, allowing to reach the maxi-mum possible resolution at ground-based telescopes with large diameters to produce diffraction limited imaging.
Fig. 6. Desired sparse pupil (left) and the aperture mask used to generate it (right). The required pupil is shown as the set of black spots and is superimposed on a scaled version of the segmented 10m Keck pri-mary mirror (hexagonal segments). From Tuthil et al. (2000).
The masks are usually categorised either as non-redundant or partially redundant. Non-redundant masks consist of arrays of small holes where two pairs of holes have not the same sepa-ration vector (the same baseline). Each pair of holes provides a set of fringes at a unique spatial frequency in the image plane.
Partially redundant masks are usually de-signed to provide a compromise between minimis-ing the redundancy of spacminimis-ings and maximisminimis-ing both the throughput and the range of spatial frequen-cies investigated. Although the signal-to-noise of speckle masking observations at high light level can be improved with aperture masks, the faintest limit-ing magnitude cannot be significantly improved for photon-noise limited detectors so that the principal limitation of the technique is that it is limited to relatively bright astronomical objects.
4.3. Aperture Synthesis
The direct closure phase measurements to-gether with the measurements of visibility amplitude allow one to reconstruct an image of any object us-ing three or more independent telescopes. This tech-nique has been successfully demonstrated by Bald-win et al. (1996) in the visible band at the Cam-bridge optical aperture-synthesis telescope (COAST, see Section 5.3). Each pair of telescopes in an array yields a measure of the amplitude of the spatial co-herence function of the object at a spatial frequency ~b/λwhich corresponds to one point in (u, v) plane.
In order to make an image from an interferom-eter, one needs estimates of the complex visibilities over a large portion of the (u, v) plane, both the am-plitudes and phases. The (u, v) points corresponding to a snapshot projection of the baseline are sampled from each available baseline. It is also useful to use a technique called ”super-synthesis”: the (u, v) plane is swept during an observation lasting several hours, due to Earth rotation. After a large variation of hour angle, several visibility moduli are, therefore,
measu-Fig. 7. uv-plane coverage for the observations of Baldwin et al. (1996). Each point corresponds to an individual visibility measurement.
red at different positions in (u, v) plane as presented in Fig. 7.
Reconstruction of complex images involves the knowledge of complex visibilities but the phase of a visibility may be deduced from closure-phase that has been applied in optical interferometry. The pro-cess after data acquisition consists of phase calibra-tion and visibility phase reconstruccalibra-tion from closure-phase terms by a technique similar to bispectrum processing, so that at least three configurations of three telescope interferometer is needed. Interferom-eters with two apertures have limited possibilities for image reconstructions due to absence of the phase visibility recovering but in that case different spec-tral channels can be employed for differential visibil-ity measurements (between the continuum and tral line) and use of data from one part of the spec-trum such as the continuum to calibrate another part (Jankov et al. 2001, 2002, 2003).
4.4. Nulling Interferometry
This technique, the nulling interferometry, holds a great promise for the detection and char-acterization of Earth-like extrasolar planets at mid-infrared wavelengths, because the light from the parent star arriving on axis is completely rejected (Bracewell 1978). The basic premise of nulling in-terferometry is conceptually quite simple: combine the starlight that arrives at a pair of telescopes so that at the centre of the image the two waves are exactly ±π out of phase producing a dark central fringe and effectively making the central star disap-pear. A deep destructive fringe is to be placed across the star, but light coming from sources that are off-set slightly from it (i.e. from planets) gets added rather that subtracted so that even though the star is completely blanked out, planets at the right loca-tions (near a constructive fringe) are not attenuated (Fig. 8).
Fig. 8. Nulling Interferometry Principle. The starlight that arrives at a pair of telescopes so that a deep destructive fringe is placed across the star (mak-ing it disappear), but the light com(mak-ing from a planet gets added rather that subtracted.
this concept is to introduce a±πphase shift over the central part of the point spread function produced by the densified pupil; a careful adjustment of the size of the phase mask together with pupil apodization in a coronographic setup can in principle provide an excellent nulling performance.
Nulling interometry can be performed from the Earth, especially when imaging Jupiter sized planets, but the best place for viewing earth-sized planets close to their suns is from a space-based plat-form. Both NASA and ESA have plans to launch infrared nulling interferometry telescopes into space over the next decades (see Section 7).
5. PAST INTERFEROMETERS
In this section, only the most important inter-ferometers that produced significant results in the past are presented:
5.1. I2T
The I2T (Interf´erom`etre `a 2 T´elescopes), Ob-servatoire Cˆote D’Azur, Plateau de Calern, France (Fig. 9) is the interferometer where the first direct interference fringes between separate telescopes were observed and reported by Labeyrie (1975).
Fig. 9. The I2T (Interf´erom`etre `a 2 T´elescopes), Plateau de Calern, France.
I2T made observations in the visible (e.g., Blazit et al. 1977, Thom et al. 1986) and pioneered interferometry at near-IR wavelengths (di Benedetto and Conti 1983, di Benedetto 1985).
5.2. MARK I, II and III
The Massachusetts Institute of Technology and the Naval Research Laboratory built and op-erated a series of prototype interferometers, named the Mark I, Mark II, and Mark III located on Mt. Wilson, California, US.
Active fringe tracking was first demonstrated in 1979 by the Mark I interferometer (Shao and Staelin 1980), while the most popular architecture today for the moving delay line is based on the so-lution implemented by the Mark III interferometer (Shao et al. 1988).
The Mark III was specifically designed to per-form wide-angle astrometry, but a variable baseline that could be configured from 3m to 31.5m provided the flexibility needed for a variety of astronomical programs. Because a full computer control of the siderostats and delay lines allowed automated acqui-sition of stars and data acquiacqui-sition, it could observe up to 200 stars in a single night, until 1993, when the interferometer stopped the operation.
5.3. COAST
The COAST (Cambridge Optical Aperture Synthesis Telescope) operated from 1991 until 2005 at Cambridge University, England. It was planned as a coherent array of four telescopes operating in the red and near infra-red, using Michelson interfer-ometry on baselines of up to 100m to give images with a resolution down to 1mas.
The number of telescopes (40cm apertures) has been brought up to five, and light from four of them could be interfered simultaneously. Switch-ing between two four-telescope arrays allowed takSwitch-ing data on nine baselines and seven closure triangles during a single night. Observations could be carried out in the visible (R,I) or near-infrared (J,H bands). It was the first instrument of its kind to exploit the techniques of aperture synthesis and closure phase at optical or infra-red wavelengths, producing the first images from an optical aperture synthesis array (Baldwin et al. 1996).
5.4. IOTA
Fig. 10. The IOTA (Infrared Optical Telescope Array) on Mt. Hopkins, US.
5.5. GI2T
The GI2T (Grand Interf´erom`etre `a 2 T´elescopes), Observatoire Cˆote D’Azur, Plateau de Calern, France (Mourard et al. 1994, see Fig. 11) was a successor of the I2T since 1985 and used two ”boule” telescopes with 1.5m apertures on a north-south baseline that could be reconfigured from 10m to 65m.
Fig. 11. GI2T (Grand Interf´erom`etre `a 2 T´elesco-pes), Plateau de Calern, France.
The GI2T had the capability of performing spectrally resolved interferometry. After cloture in 2006, it’s beam combination table, including a ver-satile visible spectrograph and an IR focus has been moved to CHARA array (Mourard et al. 2009).
5.6. MIRA
The MIRA (Mitake Infrared Array), National Astronomical Observatory, Japan, was an ambitious plan to build a series of interferometers with in-creasing capabilities. The first phase of this project (MIRA-I), which consisted of two 25cm telescopes with coud´e optics on a 4m baseline in Tokyo has suc-cessfully been completed in 1998, and has acquired
stellar fringes. The next step was to built an instru-ment with a slightly larger siderostats and a 30m baseline (MIRA-I.2). The interferometer ceased op-eration in 2007.
5.7. PTI
The PTI (Palomar Testbed Interferometer), NASA JPL, Mt Palomar, CA, US (Colavita et al. 1999, see Fig. 12) was built in 1995 and it consisted of three 40cm siderostats, up to 110m baselines which provided 3mas resolution in the near-infrared (H and K bands). It was developed primarily to demonstrate the utility of ground-based differential astrometry in the search for planets around nearby stars, and to develop key technologies for the Keck Interferometer and space-based missions.
Fig. 12. PTI aerial figure. PTI is located atop Palomar Mountain, US, next to the large white dome of the historic 5m Hale Telescope.
PTI was notable for being equipped with a ”dual-star” tracking system, the first of its kind, which simultaneously tracked interference fringes from a target star and a reference star against which the target was measured. This allowed to cancel some of the atmospheric effects of astronomical see-ing and to make very high precision measurements possible.
Aside from its role for the technical develop-ment of high precision ”dual-star” astrometry, the PTI was used mainly for stellar diameter measure-ments, stellar masses and binary star work. The in-strument concluded operations in 2009.
6. EXISTING INTERFEROMETERS
6.1. NPOI
the visible with 32 spectral channels covering the wavelength range from 450nm to 850nm. The imag-ing subarray consists of six movable siderostats with baseline lengths from 2.0m to 437m. The array ge-ometry has been optimized for baseline bootstrap-ping to facilitate the imaging of stellar surface struc-ture. Like the Mark III, the NPOI uses vacuum delay lines for pathlength compensation. The four-element astrometric subarray of the NPOI includes an exten-sive site metrology system that monitors the motions of the siderostats with respect to one another in order to perform wide-angle astrometry with more than 2mas precision.
Fig. 13. The NPOI (Navy Prototype Optical Inter-ferometer), located on Anderson Mesa, US.
6.2. SUSI
The SUSI (Sydney University Stellar Inter-ferometer) is an array of 13 telescopes with 14cm apertures, operating since 1991 at Sydney University Narrabri, Australia (Davis et al. 1999, see Fig. 14). The SUSI has two beam-combining systems: The original ”blue” system was designed to operate in the wavelength range 400-540 nm and employs nar-row bandwidths (typically a few nanometers). The newer ”red” system is designed to work in the range 500-950 nm. SUSI makes observations on a single baseline selected from a set of fixed north-south base-lines with lengths ranging from 5m to 640m and it can achieve angular resolutions from 20mas to 50µas.
Fig. 14. The SUSI (Sydney University Stellar Interferometer), Narrabri, Australia.
The 640m baseline length, the longest of all instruments currently operational or under construc-tion, has been chosen to resolve a sample of O stars at a wavelength of 450 nm. The focus of SUSI’s obser-vations is on improving our understanding of stellar astrophysics including: single stars (measuring ef-fective temperatures, radii and luminosities), binary stars, (as for single stars, plus measuring distances and masses), variable stars (e.g. Cepheids and Mi-ras), and emission line stars (e.g. Be and Wolf-Rayet stars).
6.3. ISI
The ISI (Infrared Spatial Interferometer), University of California at Berkeley, US is located close to the CHARA array and the former site of the Mark III on Mt. Wilson. It started operation in 1990 as two-telescope interferometer but presently consists of three 1.65m telescopes observing in the mid-infrared. The telescopes are fully mobile and their current site on Mount Wilson allows for place-ments as far as 70m apart providing a resolution of 3mas at 11µm. On July 2003, the ISI recorded it’s first closure phase aperture synthesis measurements. The Fig. 15 shows three ISI’s telescopes all in a line which were used for initial testing purposes, with 4m, 8m, and 12m baselines. However, they can be moved in such a way that they form a triangle and three baselines at three different angles can be measured simultaneously providing the closure phase measurments.
Fig. 15. The ISI (Infrared Spatial Interferometer), Mt. Wilson, US.
The interferometer operates at wavelengths between 9µm and 12µm and the stellar radiation is mixed with the output of a CO2 laser, which acts as the local oscillator. Observations have been carried out with baseline lengths up to 56m. One of many advantages of the ISI’s narrow heterodyne detection bandwidth is that one can tune the detection wave-length to be deliberately in or out of known spectral lines, allowing for interferometry on spectral lines to be carried out.
for making very precise measurements of positions of stars and particularly to measure positions of in-frared stars which are hardly detectable in visible light. Investigations in the field of astrometry should help to tie the astronomical reference frames together with that established in the radio, infrared and visi-ble.
6.4. CHARA
The CHARA array, named after Center for High Angular Resolution Astronomy, Georgia State University is located on Mt Wilson, US (see Fig. 16). It consists of 6 telescopes of 1m in diameter arranged in a Y-shaped configuration with baselines ranging from 30m to 330m. First fringes on a sin-gle baseline have been obtained in November 1999 and commissioning of the full array continues. Light from the individual telescopes is conveyed through vacuum tubes to the central Beam Synthesis Facil-ity in which the six beams can be combined together. When the paths of the individual beams are matched to an accuracy of less than one micron, after the light traverses distances of hundreds of meters, the Array then acts like a single coherent telescope for the pur-poses of achieving an exceptionally high angular res-olution. The Array is capable of resolving details as small as 200µas.
Fig. 16. The facilities of the CHARA (Center for High Angular Resolution Astronomy) Array to-gether with the existing facilities of historic Mount Wilson Observatory. The Beam Synthesis Facility, an L-shaped structure of the length of a football field, can be seen close to the 2.5m telescope dome.
CHARA has entered into several collabora-tions with groups offering unique instruments or technologies for enhanced performance. These in-ternational collaborations have brought a significant added value to the science capabilities of the CHARA Array. At present these collaborations include a joint observing collaboration with the Observatoire de Paris through the FLUOR instrument which has been moved from the IOTA interferometer (Section 5.4) and upgraded for the CHARA Beam Synthe-sis Facility. A collaboration with the University of
Michigan has led to the development of an imaging beam combiner (MIRC) which has already produced the first images of stellar surfaces and close binary stars. A joint project with the University of Syd-ney has led to the development of a second beam combiner with significantly improved sensitivity to fainter objects while also providing measurements of very high precision. Finally, an agreement with Ob-servatoire de la Cˆote d’Azur has brought the third new beam combiner (VEGA) to the Array capable of providing spectroscopic and polarimetric channels for high resolution work.
The first four-telescope fringes with VEGA Beam Combiner at the CHARA Array were obtained with CHARA (MIRC used as the infrared fringe tracker) in October 2010.
6.5. KI
The KI (Keck Interferometer), NASA JPL, on Mauna Kea, Hawai, US (Fig. 17) which consists of two 10m Keck telescopes obtained first fringes with full-apertures on March 2001.
Fig. 17. Two 10m telescopes of the KI (Keck In-terferometer) on Mauna Kea, Hawai, US.
The 10m telescopes are equipped with high-order adaptive optics systems, which provide good correction in the near-IR and excellent wavefront quality at 10µm. It enables the implementation of a nulling beam combiner, which will be used to char-acterize exo-zodiacal emission around nearby main-sequence stars. The standard science modes of the instrument are:
1. V2 mode, high sensitivity visibility ampli-tude measurements in the near infrared (H and K).
2. Nulling interferometry in Nband (8 -12µm) to suppress light from central star and to reveal faint extended emission around it. Key sci-entific goals are to characterize exo-zodiacal emis-sion around Sun-like stars to inform Terestial Planet Finder (see Section 7.3) mission design.
Recently, the Keck interferometer was up-graded to do self-phase-referencing (SPR) assisted K-band spectroscopy at R ≈2000. This SPR mode as currently being implemented as the ASTrometric and phase-Referenced Astronomy (ASTRA) project will provide phase referencing and astrometric ob-servations at the Keck Interferometer, leading to en-hanced sensitivity, and the ability to monitor orbits at an accuracy level of 30-100µas, high spectral res-olution observation of Young Stellar Objects (YSO), sensitive Active Galactic Nuclei (AGN) observations, and astrometric (known exo-planetary systems char-acterization and galactic center general relativity) capabilities. With the first high spectral resolution mode now offered to the community, this contribu-tion focuses on the progress of the dual field and astrometric modes.
It is planned that the KI will be upgraded with four new 1.8m ”outrigger” telescopes in order to make possible imaging mode. A large fraction of the observing time available with the outriggers should be devoted to an astrometric search for extrasolar planets while the combination of all six telescopes should result in a sensitive imaging array.
6.6. VLTI
The VLTI (Very Large Telescope Interferome-ter), European Southern Observatory Paranal, Chile, obtained first fringes on the sky in March 2001, with siderostats, and presently it allows the interferomet-ric combination of the four 8.2m telescopes (Fig. 18), augmented by four movable 1.8m auxiliary telescopes (Fig. 19).
The first generation of instruments included a commissioning instrument based on a two-telescope near-IR beam combiner VINCI (now decommis-sioned) as well as actually operating instruments MIDI, AMBER and PRIMA:
MIDI (MID-infrared Interferometric instru-ment) is the mid-infrared (N-band = 8 to 13µm) combiner which combines two beams, either from the VLT main 8.2m Unit Telescopes (UTs) or from the
Fig. 18. Four 8.2m main Unit Telescopes of the VLTI (Very Large Telescope Interferometer), Paranal, Chile.
Fig. 19. Four 1.8m Auxilary Telescopes of the VLTI (Very Large Telescope Interferometer), Paranal, Chile.
1.8m Auxiliary Telescopes (ATs), to provide visibil-ity moduli in the (u, v) plane. MIDI features spectro-scopic optics to provide visibilities at different wave-lengths within the N band and with spectral resolu-tion R=30 with the prism, R=230 with the grism (a combination of a diffraction grating and prism). UT adaptive optics guarantees diffraction-limited image quality on MIDI, for targets brighter thanV = 17m. These unique capabilities allowed an access to vari-ous observing programs including observation of the Galactic Center, analysis of the nuclei of the AGN galaxies, particularly stellar and circumstellar ob-jects, as well as solar system objects.
AMBER is the near-infrared instrument for the VLTI, which offers the possibility of combining two or three beams from either the UTs or the ATs. With spectral resolution up to 10 000, high visibility accuracy and the ability to obtain closure phases, AMBER offers the means to perform high quality interferometric measurements in the J,H,K bands (1-2.5µm range). Angular resolution is set by the max-imum available baseline, which is about 200 meters for the ATs and about 130 meters for the UTs. Ac-cordingly, the limit is about 2mas for the ATs, and about 3mas for the UTs, in the K band. These de-sign characteristics, coupled to the VLT interferome-ter potential, open up the access to investigations of several classes of objects, from stellar to extragalac-tic astronomy including, for instance, the solar sys-tem objects, low-mass stars, Cepheids, microlensing, hot exo-planets, the galactic center, nearby galax-ies, and Young Stellar Objects, Luminous Blue Vari-ables, Be stars, Novae, Wolf-Rayet and Post-AGB stars for which the results have been already pub-lished.
PRIMA can be used either to measure the angu-lar separation between the two objects (astrometry mode) or to produce images of the fainter of the two objects using a phase reference technique (imaging mode). With the limiting magnitude K≈19m with the UTs andK≈15mwith the ATs and the angular distance between the targets that can be evaluated with a resolution of 10µas, the objects such as extra-solar planets, galactic center, Active Galactic Nuclei and extragalactic objects as well as gravitational micro-lensing are the targets for PRIMA.
In addition to these ESO’s instruments there is a possibility to bring and install the visiting instru-ments. One of the visiting instruments PIONIER, developed at LAOG in Grenoble, complements ex-isting AMBER and MIDI instruments. While AM-BER has previously combined the light from three of the telescopes at the VLTI, the light coming from the four 1.8-metre Auxiliary Telescopes has been suc-cessfully combined for the first time, in October 2010. This is an important step towards unleashing the full potential of the VLTI to use multiple telescopes to-gether.
The second generation of VLTI instruments GRAVITY, VSI and MATISSE are under develop-ment:
GRAVITY will coherently combine the light from all four UTs. It will provide near infrared adap-tive optics assisted precision narrow-angle astrom-etry (10µas) and interferometric phase referenced imaging (4mas) of faint objects (K=20m). It will precisely measure the fringe visibility and phase of the fringe pattern to acquire spatial knowledge of the observed object. With an accuracy of 10µas, GRAV-ITY will be able to study motions to within a few times the event horizon size of the Galactic Center massive black hole and potentially test the General Relativity in its strong field limit. It will be able to detect intermediate mass black holes throughout the Galaxy by their gravitational action on surrounding stars. It will directly determine masses of exo-planets and brown dwarfs, and trace the origin of protostel-lar jets. Also it will be capable of spatially resolving coherent gas motions in the broad line regions of Ac-tive Galactic Nuclei in external galaxies. Through its high performance infrared wavefront sensing system, GRAVITY will open up deep interferometric imag-ing studies of stellar and gas components in dusty, obscured regions, such as obscured active galactic nuclei, dust-embedded star forming regions, and pro-toplanetary disks.
VSI (V(LTI) Spectro-Imager) was proposed as a second-generation instrument in order to pro-vide the community with spectrally-resolved, near-infrared images at angular resolutions down to 1.1mas and spectral resolutions up to R = 12 000. Targets as faint as K= 13m will be imaged without requiring a brighter nearby reference object while fainter targets can be accessed if a suitable refer-ence is available. The unique combination of high-dynamic-range imaging at high angular resolution and high spectral resolution enables a scientific pro-gram which serves a broad user community and at the same time provides the opportunity for break-throughs in many areas of astrophysics. The high level specifications of the instrument are derived from
a detailed science case based on the capability to obtain milli-arcsecond resolution images of a wide range of targets including: probing the initial condi-tions for planet formation in the AU-scale environ-ments of young stars; imaging convective cells and other phenomena on the surfaces of stars; mapping the chemical and physical environments of evolved stars, stellar remnants, and stellar winds, and disen-tangling the central regions of active galactic nuclei and supermassive black holes. The VSI will provide these new capabilities using technologies which have been extensively tested in the past. It is quite flexi-ble instrument and should be aflexi-ble to make maximum use of new infrastructure as it becomes available; for example, by combining up to 8 telescopes, enabling rapid imaging through the measurement of up to 28 visibilities in every wavelength channel within a few minutes. The current studies are focused on a 4-telescope version with an upgrade to a 6-4-telescope one.
MATISSE (Multi AperTure mid-Infrared SpectroScopic Experiment) is foreseen as a mid-infrared spectro-interferometer combining up to four UTs or ATs beams. It will measure closure phase re-lations and thus offer an efficient capability for image reconstruction. In addition to the N band, the MA-TISSE will open three new observing windows at the VLTI: the L, M, and Q bands, all belonging to the mid-infrared domain. Furthermore, the instrument will offer the possibility to perform simultaneous ob-servations in separate bands while providing several spectroscopic modes. The unique performance of the instrument is related to the existence of the four UTs that permits to push the sensitivity limits to values required by selected astrophysical programs such as the study of Active Galactic Nuclei and protoplane-tary discs. Moreover, the evaluated performance of MATISSE is linked to the availability of ATs which are relocatable in position in about 30 different sta-tions allowing the exploration of the Fourier plane with up to 200 meters baseline length. Key science programs using the ATs cover for example the forma-tion and evoluforma-tion of planetary systems, the birth of massive stars as well as the observation of the high-contrast environment of hot and evolved stars. In summary, MATISSE can be seen as a successor of MIDI by providing imaging capabilities in the en-tire mid-infrared accessible from the ground. The extension of MATISSE down to 2.7µm as well as its generalization of the use of closure phases makes it also a successor of AMBER.
6.7. LBTI
The LBTI (Large Binocular Telescope Inter-ferometer) on Mt. Graham, Arizona, US (Fig. 20) consists of two 8.4m telescopes mounted side by side in a single alt-azimuth mount, with a 14.4m center-to-center spacing.
Fig. 20. The LBT (Large Binocular Telescope) on Mt. Graham, US.
is designed for high spatial resolution, high dynamic range imaging and nulling interferometry. It con-sists of an universal beam combiner (UBC) and three camera ports, one of which is populated currently by the Nulling and Imaging Camera (NIC).
One of the key projects for the LBTI is to sur-vey nearby stars for debris disks down to levels which may obscure detection of Earth-like planets. A sur-vey of approximately 80 nearby stars that is essential as a ground-based precursor of the space missions (as for example TPF/Darwin, see Sections 7.3 and 7.4). A survey will identify zodiacal disks, which signal the existence of a planetary system for follow up study with TPF/Darwin. On the other hand, the survey will also weed out those systems with very intense zo-diacal emission, which can overwhelm the signal of terrestrial planets. Finally, the survey can also ad-dress the question of zodiacal emission strength as a function of spectral type and age of the central star. In addition, the consortium of German insti-tutes and observatories led by the Max Planck Insti-tute for Astronomy in Heidelberg together with the Istituto Nazionale di Astrofisica in Arcetri, is devel-oping the LINC ((L)BT (IN)terferometric (C)amera) instrument. The LINC will combine the radiation from the two 8.4m primary mirrors in Fizeau mode which allow true imagery over a wide field of view. The beam combiner will operate at wavelengths be-tween 0.6 and 2.4µm. When coupled with the ad-vanced adaptive optics system of the LBT, the LINC instrument will allow the interferometric measure-ments over a field of approximately 10-20 arcseconds square. Ultimately, after improvement of the Adap-tive Optics, which will allow better performance over a wider field of view, the instrument will be known as NIRVANA, the (N)ear-(IR) / (V)isible (A)daptive I(N)terferometer for (A)stronomy. The main scien-tific goals of LINC/NIRVANA are YSOs environ-ments, circumstellar disks and outflows, formation of binary stars, stellar clusters, the compact H II regions, astrometry for extra-solar planets, the so-lar system minor bodies, the galactic center and the galaxies at z≈1−2.
In October 2010 the LBTI acquired its first fringes on the sky and will be interferometrically functional in 2011 once the adaptive optics and the LBTI are commissioned.
6.8. MROI
The MROI (Magdalena Ridge Observatory In-terferometer) is presently under construction and should be operational in years to come. It is an inter-national scientific collaboration between New Mexico Tech (US) and the University of Cambridge (UK). The Interferometer Project’s mission is to develop a ten-element imaging interferometer to operate at wavelengths between 0.6µm and 2.4µm with base-lines from 7.5m to 340m. Its primarily technical and scientific goals are to produce model-independent im-ages of faint and complex astronomical targets at resolutions over 100 times that of the Hubble Space Telescope in order to study: star and planet for-mation, stellar accretion and mass loss, and active galactic nuclei.
The 2.4-meter telescope achieved first light on October 2006 and commenced operations on Septem-ber 2008. The work at the interferometer site started in September 2006 and the Beam Combining Facil-ity was completed in early 2008. Presently, the tele-scopes are nearing their first factory acceptance tests. When finished, the interferometer will be upgraded with ten telescopes, each approximately 1.4 meters in diameter. The telescopes making up the interfer-ometer will be spaced by distances up to 400 meters and will be optically linked in order to make aper-ture synthesis imaging. This setup will simulate the resolving power of a single telescope up to 400 me-ters in diameter. As a result of the large number of telescopes in the array, the interferometer will be able to make accurate images of complex astronom-ical objects.
7. FUTURE INTERFEROMETERS 7.1. OHANA
The OHANA (Optical Hawaiian Array for Nanoradian Astronomy) is the project of a hecto-metric fibered array at Mauna Kea, Hawaii. The concentration of large apertures on Mauna Kea (two Keck, Gemini, Subaru, CFHT, UKIRT) lends itself to ideas for interferometric combination of these six telescopes (Mariotti et al. 1996). By taking advan-tage of the existing adaptive optics systems, such as already installed on telescopes, this would provide excellent sensitivity and unprecedented imaging ca-pabilities with baselines ranging from 75m to 800m (the largest baseline of OHANA (Subaru-Gemini)). This yields resolutions of 0.25mas and 0.5mas at 1 and 2 microns respectively, therefore opening the possibilities for research in many fields as for VLTI, but with four times better spatial resolution.
interfer-ometer on July 2010, the first step in a plan to link the giant telescopes of Hawaii into a powerful optical interferometer.
7.2. Interferometry at Antartica
The Dome C at Antartica (altitude 3300m) is considered as a perfect site for astronomical inter-ferometry due to low wind speeds, little atmospheric turbulence, negligible seismic activity and very sta-ble climate. For this reason, several interferometric projects are under development, most notably the API (Antartic Planet Interferometer, Swain et al. 2003) and KEOPS (Kiloparsec Explorer for Optical Planet Search, Vakili et al. 2003).
TheAPIshould be built in two phases. Phase 1 is proposed to consist of a two-element active fringe tracking interferometer using two 50cm diameter telescopes operating in a direct visibility (V2) mode with baselines of up to 400 metres. The H,K,L,M bands (1.4 to 5.4µm) are to be targeted to give max-imum resolution and sensitivity at up to 104spectral resolution. In Phase 2, the full API, will consist of several 2m diameter telescopes and will be capable of operating in a variety of modes. It will be a partic-ularly powerful tool for the study of extrasolar plan-ets, protoplanplan-ets, Young Stellar Objects, and Active Galactic Nuclei accretion disks.
The KEOPS is proposed as three concentric rings of 1 to 2m diameter telescopes covering a di-ameter of 650 metres. The 36 co-phased primary telescopes would have an equivalent collecting area of at least 60m2and a resolution of a milli-arcsecond at 10µm. KEOPS is designed for direct imaging of exo-planets in the thermal infrared, as well as for astroseismology studies.
7.3. TPF
The TPF (Terrestrial Planet Finder) Interfer-ometer is a concept for a formation flying interferom-eter to search for Earth-like planets around nearby stars, planet studies and signs of life. TPF was be-ing developed as a possible collaboration with ESA’s Darwin mission but the project was recently can-celled.
The NASA’s 2010 Decadal Survey recom-mended a continuation of technology development for nulling interferometry under the title of ”New Worlds Technology Development Program” and now, although TPF no longer exists as a planned mis-sion within NASA, some technology work continues. Efforts in 2009 and 2010 have included broadband nulling with the Adaptive Nuller and demonstrations with the Planet Detection Testbed. One of the pos-sible configurations of a space interferometer (an op-tion considered for Terrestial Planet Finder) is pre-sented in Fig. 21.
At present, NASA continues to work with Uni-versity scientists and industry to develop the TPF technology. Meanwhile, the Kepler mission will col-lect information on the number of possible nearby Earthlike planets, while the Keck Interferometer and Large Binocular Telescope Interferometer projects will measure, for this purpose, the dust surrounding those stars.
Fig. 21. A possible configuration of a space inter-ferometer.
7.4. Darwin
The Darwin mission was planned as a con-stellation of four or five spacecrafts, with primary goal to look for Earth-like planets and analyse their atmospheres for chemical signatures of life. One spacecraft would have acted as a central communica-tions station. The other three/four would have func-tioned as light collectors, redirecting light beams to the central station. The constellation was also in-tended to carry out high-resolution imaging using aperture synthesis, to provide images of celestial ob-jects with unprecedented detail. Instead of orbit-ing Earth, Darwin was proposed to be placed be-yond the Moon. At a distance of 1.5 million km from Earth, in a direction opposite to that of the Sun, Darwin would have operated from the second Lagrange point L2. The Darwin mission was pro-posed in 2007 as a concept for Phase A Study within ESA’s Cosmic Vision, which includes missions be-yond 2015. It was not chosen for further study, but a similar project should be re-proposed to ESA. What-ever the final design of the telescopes, it is clear that the TPF/Darwin technology will be a pathfinder for future micro and nano-arcsecond resolution instru-ments at infrared and other wavelengths.
8. CONCLUSION
After decades of development, optical interfer-ometry is now beginning to play a major role in astro-physics. The emergence of large interferometers, in particular the Very Large Telescope interferometer, CHARA, the Keck Interferometer and Large Binocu-lar Telescope Interferometer promise to revolutionize the impact of high resolution observations in many areas of astrophysics. Current programs execute in-creasingly varied and sophisticated observations. Ini-tially, these were almost exclusively observations of stars and circumstellar phenomena, but recently, op-tical interferometry extended even to extragalactic targets.
galactic astronomy, in particular the areas of fun-damental stellar properties, star and planet forma-tion, and all stages of stellar evoluforma-tion, but increas-ingly, optical interferometry will play more impor-tant role in other areas of astrophysics particularly exo-planetary and extragalactic research.
Acknowledgements – This work has been supported by the Ministry of Science and Technological Devel-opment of the Republic of Serbia (Project No 146007 ”Inverse problems in astrophysics : Interferometry and Spectrophotometry of stars”).
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ASTRONOMSKA OPTIQKA INTERFEROMETRIJA.
I METODE I INSTRUMENTI
S. Jankov
Astronomical Observatory Belgrade, Volgina 7, 11060 Belgrade 38, Serbia
E–mail: [email protected]
UDK 520.36–13
Pregledni rad po pozivu
Prethodna dekada je videla ostva-renje velikih interferometrijskih projeka-ta ukljuqujui teleskope klase 8-10m i stometarske interferometrijske baze. Mo-derni kompjuteri i kontrolna tehnologija su omoguili interferometrijsku kombi-naciju svetlosti sa razdvojenih teleskopa kako u vidljivom tako i u infracrvenom reжimu. Oslikavanje sa rezolucijom od mili-luqne sekunde i astrometrija sa preciznoxu od mikro-luqne sekunde su tako postale real-nost. Ovde dajem pregled metoda i instru-mentacije koje odgovaraju tekuem stanju u oblasti astronomske optiqke interferomet-rije. Prvo, ovaj pregled sumarizuje razvoj od pionirskih radova Fizoa i Majkelsona. Za-tim su opisane osnovne dostupne veliqine a posle toga sledi diskusija o osnovnim
